2020-12-21T10:43:46Z
Acceptance in OA@INAF
þÿOverview of non-transient ³-ray binaries and prospects for the Cherenkov
Telescope Array
Title
Chernyakova, M.; Malyshev, D.; PAIZIS, ADAMANTIA; LA PALOMBARA, NICOLA;
Balbo, M.; et al.
Authors
10.1051/0004-6361/201936501
DOI
http://hdl.handle.net/20.500.12386/29039
Handle
ASTRONOMY & ASTROPHYSICS
Journal
631
A&A 631, A177 (2019) https://doi.org/10.1051/0004-6361/201936501 c ESO 2019
Astronomy
&
Astrophysics
Overview of non-transient
γ
-ray binaries and prospects
for the Cherenkov Telescope Array
M. Chernyakova
1,2, D. Malyshev
3, A. Paizis
4, N. La Palombara
4, M. Balbo
9, R. Walter
9, B. Hnatyk
10,
B. van Soelen
11, P. Romano
5, P. Munar-Adrover
6, Ie. Vovk
7, G. Piano
8, F. Capitanio
8, D. Falceta-Gonçalves
16,
M. Landoni
5, P. L. Luque-Escamilla
15, J. Martí
15, J. M. Paredes
12, M. Ribó
12, S. Safi-Harb
13, L. Saha
14,
L. Sidoli
4, and S. Vercellone
51 School of Physical Sciences and CfAR, Dublin City University, Dublin 9, Ireland
e-mail: masha.chernyakova@dcu.ie
2 Dublin Institute for Advanced Studies, 31 Fitzwilliam Place, Dublin 2, Ireland
3 Institut für Astronomie und Astrophysik Tübingen, Universität Tübingen, Sand 1, 72076 Tübingen, Germany 4 INAF – IASF Milano, Via Alfonso Corti 12, 20133 Milano, Italy
5 INAF – Osservatorio Astronomico di Brera, Via E. Bianchi 46, 23807 Merate, Italy
6 Unitat de Física de les Radiacions, Departament de Física, and CERES-IEEC, Universitat Autònoma de Barcelona,
08193 Bellaterra, Spain
7 Max-Planck-Institut für Physik, 80805 München, Germany 8 INAF – IAPS Roma, Italy
9 ISDC, University of Geneva, Switzerland
10 Astronomical Observatory of Taras Shevchenko National University of Kyiv, 3 Observatorna str. Kyiv, 04053, Ukraine 11 University of the Free State Department of Physics, PO Box 339, 9300 Bloemfontein, South Africa
12 Institut de Ciències del Cosmos (ICCUB), Universitat de Barcelona, IEEC-UB, Martí i Franquès 1, 08028 Barcelona, Spain 13 Dept of Physics and Astronomy, University of Manitoba, Winnipeg R3T 2N2, Canada
14 Unidad de Partículas y Cosmología (UPARCOS), Universidad Complutense, 28040 Madrid, Spain 15 EPS Jaén, Universidad de Jaén, Campus Las Lagunillas s/n, 23071 Jaén, Spain
16 Escola de Artes, Ciências e Humanidades, Universidade de São Paulo, Rua Arlindo Bettio 1000, 03828-000 São Paulo, Brazil
Received 13 August 2019/ Accepted 22 September 2019
ABSTRACT
Aims.Despite recent progress in the field, there are still many open questions regarding γ-ray binaries. In this paper we provide an
overview of non-transient γ-ray binaries and discuss how observations with the Cherenkov Telescope Array (CTA) will contribute to their study.
Methods.We simulated the spectral behaviour of the non-transient γ-ray binaries using archival observations as a reference. With this
we tested the CTA capability to measure the spectral parameters of the sources and detect variability on various timescales.
Results.We review the known properties of γ-ray binaries and the theoretical models that have been used to describe their spectral
and timing characteristics. We show that the CTA is capable of studying these sources on timescales comparable to their characteristic variability timescales. For most of the binaries, the unprecedented sensitivity of the CTA will allow studying the spectral evolution on a timescale as short as 30 min. This will enable a direct comparison of the TeV and lower energy (radio to GeV) properties of these sources from simultaneous observations. We also review the source-specific questions that can be addressed with these high-accuracy CTA measurements.
Key words. gamma-rays: stars – gamma-rays: general – acceleration of particles – methods: observational – binaries: general
1. Introduction
γ-ray binaries are a subclass of high-mass binary systems whose energy spectrum peaks at high energies (HE, E & 100 MeV) and extends to very high energy (VHE, E & 100 GeV) γ-rays. In these systems a compact object (neutron star, NS, or a black hole, BH) orbits a young and massive star of either O or B type. While high-mass binaries represent a substantial frac-tion of the Galactic X-ray sources that are detected above 2 keV (e.g. Grimm et al. 2002), fewer than ten binaries were detected in the γ-ray band by the current generation of Cherenkov telescopes, such as Major Atmospheric Gamma Imaging Cherenkov Telescopes (MAGIC;Aleksi´c et al. 2012a), Very Energetic Radiation Imaging Telescope Array System
(VERITAS; Park & VERITAS Collaboration 2015), and High Energy Stereoscopic System (H.E.S.S.;Aharonian et al. 2006). Therefore γ-ray binaries represent a relatively new and unex-plored class of astrophysical objects.
Of all the binary systems that are regularly observed at TeV energies, the nature of compact objects is only firmly estab-lished for two systems, PSR B1259−63 and PSR J2032+4127. In PSR B1259−63, a 43 ms radio pulsar orbits a Be star on a highly eccentric 3.4 yr orbit. Radio pulsations from the source are detected along most of the orbit, except for a brief period near periastron. The second source, which also con-tains a pulsar, PSR J2032+4127, has an even longer orbital period of about 50 yr (Lyne et al. 2015; Ho et al. 2017), and TeV emission was detected from this system by VERITAS and
MAGIC as the pulsar approached its periastron in September 2017 (VERITAS & MAGIC Collaborations 2017).
All other known systems are more compact, and the nature of their compact object is so far unknown. It is possible that these systems harbour radio pulsars as well, but the optical depth due to the stellar wind outflow is too high to detect the radio signal that originates close to the pulsar (the so-called hidden pulsar model, see e.g.Zdziarski et al. 2010). Alternatively, it is possible that some of these systems harbour a BH or an accreting NS (the microquasar model,Mirabel & Rodríguez 1998).
Of the accretion-powered γ-ray binaries that likely contain a BH, the highest energy emission that has so far been regularly observed comes from Cyg X-1 and Cyg X-3 and is detected with AGILE and the Fermi Large Area Telescope (Fermi-LAT; e.g. Tavani et al. 2009a;Sabatini et al. 2010; Bodaghee et al. 2013; Malyshev et al. 2013). Cyg X-1 is detected up to about 10 GeV in the hard state (Zanin et al. 2016; Zdziarski et al. 2017), and Cyg X-3 is detected up to about 10 GeV during flares that mostly occur when the source is in the soft state (Zdziarski et al. 2018). In 2006, the MAGIC telescope has also reported marginal detec-tion of a TeV flare from Cyg X-1 at a 3.2σ confidence level. This coincides with an X-ray flare seen by RXTE, Swift, and INTEGRAL (Albert et al. 2007). We therefore currently have evidence that accreting sources can accelerate particles only dur-ing some very specific states, and we need to study the persis-tent γ-ray binaries in detail to unveil their ability to steadily accelerate particles (at least at given orbital phases). Recently, the microquasar SS433 was also detected by Fermi-LAT and HAWC (Bordas et al. 2015; Abeysekara et al. 2018). The per-sistent emission reported by HAWC is localised to structures in the lobes, far from the centre of the system. This implies a very different emission scenario to the other systems.
In addition to the γ-ray binaries that contain a compact object, HE and VHE γ-rays have also been detected from colliding-wind binaries (CWB). A CWB is a binary star system consisting of two non-compact massive stars that emit powerful stellar winds, with high mass-loss rates and high wind veloci-ties. The collision of the winds produces two strong shock fronts, one for each wind. They surround a shock region of compressed and heated plasma, where particles are accelerated to VHEs (Eichler & Usov 1993). To date, only one CWB has a confirmed detection at VHE: η Carinae (η Car).
The family of γ-ray binaries can be extended through follow-up observations in the next years based on indications from lower energy bands. In particular, binaries with pul-sars orbiting Be or O stars are likely to provide a notice-able addition to the γ-ray binary list. Similarly, BH and Be star binaries can also be considered good candidates for the γ-ray binary family (Williams et al. 2010;Munar-Adrover et al. 2016), although only one system, MWC 656, is known so far (Casares et al. 2014). Other alternative search efforts have focused on multi-wavelength cross-identification that explores the possible association of luminous early-type stars with GeV γ-ray sources (mainly) detected by Fermi-LAT (McSwain et al. 2013;Martí et al. 2017). Still,Dubus et al.(2017) recently car-ried out a synthetic population simulation and estimated that fewer than 230 systems exist inside the Milky Way.
In recent years, γ-ray binaries have already been the subject of numerous observational campaigns and theoretical studies (e.g.Dubus 2013), which strongly indicate that the high-energy emission from these systems is primarily powered by the outflow from the compact object. However, due to the limited sensitivity of the current generation of instruments, the nature of the compact object (NS or BH) and the details of the particle
acceleration, with efficiency sometimes close to the theoretical limit (e.g.Johnson et al. 2018), remain unknown in most of the systems (e.g.Paredes & Bordas 2019).
Existing data have shown that in some systems such as LS 5039 and LS I +61◦ 303, the observed HE and VHE
emission are separate components that are generated at differ-ent places (e.g.Zabalza et al. 2013, and references therein). A proper modelling of these double-component spectra requires time-resolved spectroscopy throughout the binary orbit. In addi-tion, γ-ray binaries are known to be variable on timescales as short as hours, minutes, and even tens of seconds, as observed in X-ray and HE bands (e.g.Chernyakova et al. 2009;Smith et al. 2009; Johnson et al. 2018). At the same time, with the cur-rent generation of VHE telescopes, observations can only pro-vide information that is averaged over several days even for the brightest binaries. The possibility of studying the broad-band spectral variability on a characteristic timescales is crucial for an unambiguous modelling.
In the next decade this situation may change with the deploy-ment of the next-generation VHE telescope, the Cherenkov Tele-scope Array (CTA) observatory.
The CTA will be composed of two sites, one in the Northern (La Palma, Canary Islands) and one in the Southern Hemisphere (Paranal Observatory, Chile), which will enable observations to cover the entire Galactic plane and a large fraction of the extra-galactic sky (see e.g. CTA Consortium 2017). The array will include three different telescope sizes to maximise the energy range of the instrument (from 20 GeV to more than 300 TeV). With more than 100 telescopes in the Northern and Southern Hemispheres, the CTA will be the largest ground-based γ-ray observatory in the world. The CTA will be 5–20 times more sen-sitive (depending on the energy) than the current generation of ground-based γ-ray detectors (CTA Consortium 2019). It is fore-seen that CTA will make a breakthrough in many areas, includ-ing the study of γ-ray binaries. Beyond detailed studies of the known binaries, the CTA is foreseen to discover new sources, enlarging the population.Dubus et al.(2017) has estimated that four new γ-ray binaries can be expected to be detected in the first two years of the CTA Galactic Plane survey.
The aim of this paper is to estimate the potential of the CTA for observations of known γ-ray binary systems. The text is organised as follows. In Sect.2we outline the source selection and CTA simulation setup. Sections3and4present the results of simulations for specific binary system types. Finally, in Sect.5 we briefly summarise and discuss our results.
2. Simulations
All the simulations that are reported in this paper were per-formed with the ctools analysis package1 (Knödlseder et al.
2016, v 1.5), together with the prod3b-v1 set of instrument response functions (IRFs2) for both the northern (La Palma) and
southern (Paranal) CTA sites. In prod3b-v1 IRFs only exist for zenith angles of 20 and 40◦. To select the correct response
func-tion, we used a simple relation between the minimum zenith angle of the source (mza) and declination (Dec) and the latitude (lat) of the site: mza = |lat − Dec|. For example, for La Palma (lat= +29◦), HESS J0632+057 has an mza = 23◦. Thus the 20◦
IRF is the most appropriate. For the southern site (Paranal,
1 http://cta.irap.omp.eu/ctools/
2 https://www.cta-observatory.org/science/
Table 1.Properties of γ-ray binaries with a compact source. PSR LS LS I HESS 1 FGL HESS LMC P3(∗∗) B1259−63(?) 5039(†) 61◦303(•) J0632+057() J1018.6−5856(‡) J1832−093 Porb(days) 1236.724526(6) 3.90603(8) 26.496(3) 315(5) 16.544(8) – 10.301(2) e 0.86987970(6) 0.24(8) 0.54(3) 0.83(8) 0.31(16) - 0.40(7) ω (◦) 138.665013(11)(]) 212(5) 41(6) 129(17) 89(30) – 11(12) i(◦) 153.3+3.2 −3.0 13–64 10–60 47–80 – – – d(kpc)(1) 2.39 ± 0.18 2.07 ± 0.22 2.63 ± 0.26 2.76 ± 0.34 6.52 ± 1.08 – 50.0 ± 1
Spectral type O9.5Ve O6.5V(f) B0Ve B0Vpe O6V(f) – O5 III(f)
M?(M ) 14.2–29.8 23 12 16 31 – – R?(R ) 9.2 9.3 10 8 10.1 – – T?(K) 33 500 39 000 22 500 30 000 38 900 – 40 000 dperiastron(AU) 0.94 0.09 0.19 0.40 (0.41) – – dapastron(AU) 13.4 0.19 0.64 4.35 (0.41) – – φperiastron 0 0 0.23 0.967 – – 0.13 φsup. conj. 0.995 0.080 0.036 0.063 – – 0.98 φinf. conj. 0.048 0.769 0.267 0.961 – – 0.24
IRF: South_z40 South_z20 North_z20 South_z40 South_z40 South_z20 South_z40
North_z40 North_z20 North_z40
Notes.(])Argument of periastron of the pulsar orbit (massive star for the other systems).(1)All distances given with an error are taken from the Gaiaarchive,https://gea.esac.esa.int/archive/
References.(?)Shannon et al.(2014),Miller-Jones et al.(2018),Negueruela et al.(2011).(†)Ribó et al.(2002),McSwain et al.(2004),Sarty et al. (2011).(•)McSwain et al.(2004), Aragona et al.(2009).()Casares et al. (2012), Aliu et al. (2014), Aragona et al.(2010).(‡)An et al.(2015),
Monageng et al.(2017),Napoli et al.(2011).(∗∗)Corbet et al.(2016),Pietrzy´nski et al.(2013).
lat= −25◦), HESS J0632+057 has a minimal zenith angle of 31◦,
and we chose the 40◦IRF.
In the analysis, we simulated the data with ctobssim and fit-ted simulafit-ted event files with ctlike using a maximum likelihood method. To simulate the event file, we used all sources listed in TeVCat3within a circle of 5◦around the source position. In
addi-tion, we included the instrumental background and the Galactic diffuse γ-ray emission in the model of the region surrounding the simulated source4. All errors presented in the paper are the
statistical errors at a 1σ confidence level.
3.
γ
-ray binaries with a compact sourceThis section is devoted to an overview of the non-transient, point-like, γ-ray binary sources that all consist of an O- or B/Be-type star and a compact object (pulsar or BH). The specific sources studied here are listed in Table1. PSR J2032+4127 is not included because with a ≈50 yr orbital period, it is unlikely that the next periastron passage will be observed with the CTA. Very recently, while this paper was under revision, a new γ-ray binary candidate, 4FGL J1405.1−6119, was discovered (Corbet et al. 2019). The TeV properties of the source are currently not known and we therefore do not discuss this source here either.
The VHE emission of all the γ-ray binaries is well described by a power law with an exponential high-energy cut-off. As was mentioned in the introduction, current VHE observations are not sensitive enough to follow the details of the spectral evolution of these systems on their characteristic timescales.
In order to test the future capabilities of the CTA, we cal-culated the predicted errors on the spectral parameters for dif-ferent characteristic fluxes on 30 min and 5 h timescales in the
3 http://tevcat2.uchicago.edu/
4 We have verified that results obtained have a negligible dependency
on the choice of the Galactic diffuse background model.
1–100 TeV energy range. To do this, we considered 100 ran-dom realisations of the region surrounding the binary. For each simulation we used a power-law spectral shape and assumed flux and spectral slope values that are typical for the simulated γ-ray binary. The uncertainties are defined as a standard devia-tion of the distribudevia-tion of best-fit values. The results are shown in the top and middle panels of Fig. 1. While for most sys-tems the spectral shape above 1 TeV nicely follows a power law, it is not yet clear at which energy we should expect a cut-off. In order to estimate the maximum energy up to which the CTA will be able to firmly detect a cut-off, we fitted event data (simulated for a power-law spectral energy distribution) with a cut-off power-law model. This was repeated 1000 times for dif-ferent data realisations. From the obtained distribution of best-fit off values we found a value above which 95% of all cut-off values are located. This corresponds to the 95% upper limit on the cut-off energy, that is, if the cut-off is detected by the CTA at energies lower than this, we can be confident that it is real. The resulting 95% confidence values for sources with fluxes F(>1 TeV) < 1.5 × 10−12ph cm−2s−1 are shown in the
bottom panel of Fig.1. For sources with higher fluxes, the result-ing value of the cut-off is close to 100 TeV. The spectrum of PSR B1259−63 is much softer (Γ ≈ 2.9) than those of other binaries (Γ ≈ 2.3), which results in a lower value of the cut-off energy that can be detected by the CTA.
We furthermore confirmed that for a given flux and expo-sure time, the error on the slope has a weak dependence on the slope value. Figure 2 illustrates the slope uncertainty (shown with colour) as a function of the slope and the 1–10 TeV flux level for a 5 h observation of the point-like source located at the position of PSR B1259−63.
Lastly, in this section we present an overview of what is known about each source listed in Table1. We discuss the ques-tions that can be answered with the new data measured with the precision shown in Fig.1, and present results of other simula-tions specific to each individual case.
0.00 0.05 0.10 0.15 0.20 0.25 0.30 0.35 0.40 ∆ F/F
exp=5h exp=30 min
0.00 0.05 0.10 0.15 0.20 0.25 0.30 0.35 0.40 ∆ Γ 0.0 0.5 1.0 1.5 2.0
Source flux > 1 TeV [10−12 ph/cm2/s] 100 101 102 Ecut [T eV ] PSR B1259-63HESS J0632+057 LMC P3 1FGL J1018.6-5856 LS5039 HESS J1832-093 0.0 0.1 0.2 0.3 0.4 0.5 0.6 0.7 ∆ F/F
exp=5h exp=30 min
0.00 0.05 0.10 0.15 0.20 0.25 0.30 ∆ Γ 0.0 0.5 1.0 1.5 2.0 2.5 3.0
Source flux > 1 TeV [10−12 ph/cm2/s] 100 101 102 Ecut [T eV ] LSI 61 303HESS J0632+057 LS5039 HESS J1832-093
Fig. 1.Summary of the simulations we performed for the different binaries for various exposure times and telescope configurations. Left and right
panels: south and north sites, respectively. Exposure time is shown with colours: blue corresponds to 30 min and red to 5 h. In this figure, we show the dependence of the relative flux error (top panel), spectral slope error (middle panel), and maximum energy up to which a cut-off can be excluded (bottom panel).
0.5
1.0
1.5
2.0
2.5
Integral Flux 1 10 TeV, 10
12ph/cm
2/s
4.0
3.5
3.0
2.5
2.0
Index
0.01 0.02 0.03 0.050.000
0.025
0.050
0.075
0.100
0.125
0.150
0.175
0.200
Fig. 2.Uncertainty on the slope (colour bar) as a function of the slope
value and the flux in the 1−10 TeV energy range for 5 h observations of a point-like source at the position of PSR B1259−63. White lines illustrate the levels of constant slope uncertainty.
3.1. PSR B1259−63 3.1.1. Source properties
PSR B1259−63 was first discovered as part of a search for short-period pulsars with the Parkes 64 m telescopes (Johnston et al. 1992a), and was the first radio pulsar discovered in orbit around a massive non-degenerate star, the Be star LS 2883 (Johnston et al. 1992b). Long-term monitoring of the pulsar has allowed for a very accurate determination of the binary orbit and reveals that PSR B1259−63 is on a highly eccentric 3.4 yr orbit (Shannon et al. 2014, and references therein).
Radio observations around periastron show an increase and variability in the dispersion measurement of the pulsed signal as the pulsar passes into the stellar wind (e.g. Johnston et al. 2001). This is followed by an eclipse of the pulsed signal from ≈16 days before until ≈16 days after periastron, accompanied by the detection of unpulsed radio emission (Johnston et al. 2005,
and references therein). The unpulsed emission shows a double-peak structure, reaching a maximum around the time of the start and end of the pulsar eclipse, although the shape varies between the periastron passages (e.g.Johnston et al. 2005). The unpulsed emission originates from the extended pulsar wind nebula, which is shown to extend beyond the binary by observations with the Australian Long Baseline Array (Moldón et al. 2011).
The best optical analysis of the optical companion, LS 2883, comes from high-resolution spectroscopic observations with the UVES/VLT (Negueruela et al. 2011). It is a rapidly rotating O9.5Ve star with an oblate shape and a temperature gradient from the equator to the poles. The star is wider and cooler at the equator (Req ≈ 9.7 R ; Teff,eq ≈ 27 500 K), and narrower
and hotter at the poles (Rpol ≈ 8.1 R ; Teff,pol ≈ 34 000 K).
The Be nature of the star is clear from the strong emission lines that are observed from the source, which originate from the out-flowing circumstellar disc (Johnston et al. 1992b,1994; Negueruela et al. 2011). The disc is believed to be tilted rela-tive to the orbital plane (e.g. Wex et al. 1998), with the pulsar crossing the disc plane twice per orbit. Observations have shown that the circumstellar disc is variable around periastron, with the strength of the Hα line increasing until after periastron, as well as changes in the symmetry of the double-peaked He
i
line(Chernyakova et al. 2014,2015;van Soelen et al. 2016). After first being detected at X-ray energies with ROSAT (Cominsky et al. 1994), observations around periastron have shown a remarkable similarity during different periastron pas-sages. X-ray observations folded over multiple epochs show that the X-ray flux peaks before and after periastron, at around the same time as the pulsed radio emission becomes eclipsed (e.g.Chernyakova et al. 2015, and references therein). This is interpreted as being associated with the time the pulsar passes through the plane of the circumstellar disc. Observations around the 2014 periastron passage also revealed that the rate at which the flux decreased after the second maximum (≈20 d after peri-astron) slowed down and plateaued around 30 days after perias-tron, at the time when the GeV γ-ray emission began to increase
rapidly. Extended X-ray emission has also been detected around PSR B1259−63, with an extended structure flowing away from the binary; this is suggested to be a part of the circumstellar disc that is ejected from the system and begins to become accelerated outwards by the pulsar wind (Pavlov et al. 2011,2015).
While not detected by COMPTEL and EGRET (Tavani et al. 1996), PSR B1259−63 has subsequently been detected at GeV and TeV γ-ray energies with Fermi-LAT and H.E.S.S. The H.E.S.S. telescope has reported on observations of the source over the 2004, 2007, 2010, and 2014 periastron passages (Aharonian et al. 2005a, 2009; H.E.S.S. Collaboration 2013; Romoli et al. 2017). The combined light curves over multiple epochs are beginning to show an indication of a double-hump structure around periastron, with a dip at periastron. This is sim-ilar to what is observed at X-ray energies. The observations at GeV energies with Fermi-LAT show a very different result. During the 2011 periastron passage, a very faint detection was reported around periastron, but about 30 days after periastron, a rapid brightening (flare) with a luminosity approaching that of the pulsar spin-down luminosity was observed (Tam et al. 2011; Abdo et al. 2011). This occurred at a time when the multi-wavelength emission decreased, and a flare at this period was not expected. Observations around the following periastron, in 2014, had a substantially shorter exposure before the flare and no emission was detected before or at periastron. While the flux started to increase at around the same orbital phase, the emission peaked later and was fainter than during 2011 (Tam et al. 2015; Caliandro et al. 2015). The most recent peri-astron passage in 2017 has also shown a different light curve: while no γ-ray flare was reported 26–43 days after periastron (Zhou et al. 2017), a rapid flare was detected at 70 days after periastron, during a period when GeV emission was previously not detected (Johnson et al. 2017). In addition, rapid ≈3 h flares in GeV and changing UV flux have been reported during the last periastron passage (Tam et al. 2018). Detailed analysis of the short timescale variability of the source by Johnson et al. (2018) revealed even shorter substructures on a timescale of ≈10 min. The energy released during these short flares signif-icantly exceeds the total spin-down luminosity. This demon-strates a clear variability of the emission on very different timescales from as short as few minutes up to orbit-to-orbit variability.
3.1.2. Prospects for CTA observations
The TeV γ-ray emission from PSR B1259−63 has been detected from around 100 days before until 100 days after periastron, with the next periastron occurring on 2021 February 9. The ≈3.4 yr orbital period makes observations more challenging because orbit-to-orbit variation studies must take place over long time periods. Despite this, the improved sensitivity of the CTA obser-vations around the next periastron passages can be used to test different models and better constrain the theoretical models of this source. This may include investigating the degree of gamma-gamma absorption around periastron, searching for connections to the GeV flare, and constraining the shape of the light curve near the disc crossings.
The double-hump shape of the TeV light curve around peri-astron has for instance been attributed to more efficient parti-cle acceleration during the disc crossing (Takata et al. 2012), hadronic interactions in the disc (Neronov & Chernyakova 2007), time-dependent adiabatic losses modified by the disc (Kerschhaggl 2011), and increased gamma-gamma absorption around periastron (Sushch & van Soelen 2017). Gamma-gamma
100 75 50 25 0 25 50 75 100
Days from periastron 0 2 4 6 8 Int eg ra te d En er gy Fl ux , 1 0 12 er g/ cm 2/s
Fig. 3.Simulated light curve above 1 TeV of PSR B1259−63 around
periastron. Each point has a 30 min exposure.
absorption of the TeV photons should be highest a few days before periastron, and if TeV γ-rays are produced near the pul-sar location, stellar and disc photons should decrease the flux above 1 TeV, harden the photon index, and vary the low-energy cut-off (Sushch & van Soelen 2017). The simulated light curve for PSR B1259−63 is shown in Fig.1for 30 min observations. These measurable limits will help to place better constraints on the level of γγ absorption in the system.
The second question the CTA can start to answer is whether there is any indication at TeV energies of a connection to the GeV flare. The H.E.S.S. observations around the 2010 peri-astron passage showed no TeV flare at the time of the Fermi flare (H.E.S.S. Collaboration 2013), and similarly, no multi-wavelength flare has been detected. However, X-ray observa-tions around the 2014 periastron passage showed a change in the rate at which the flux decreased (Chernyakova et al. 2015). Observations with the CTA around periastron will allow us to search for a similar effect. This will be an important constraint on the underlying emission mechanism.
Finally, the improved sensitivity of the CTA will enable a more detailed investigation of the shape of the light curve around the periods of the disc crossings. This is an impor-tant comparison to make to models that predict various shapes around these periods, such asKerschhaggl (2011) and Neronov & Chernyakova(2007).
To illustrate this point, we simulated the light curve of the source around periastron. For this we assumed a constant slope withΓ = 2.9 and modulated the flux above 1 TeV according to the H.E.S.S. observations reported by Romoli et al.(2017). Fig.3illustrates that even 30 min exposures will be enough for the CTA to measure the profile with a high accuracy.
3.2. LS I +61◦303
3.2.1. Source properties
The srouce LS I +61◦ 303 was first discovered as a bright
γ-ray source by the Cos B satellite (Hermsen et al. 1977). Shortly after the discovery, it was realised that this source was also a highly variable radio source (Gregory & Taylor 1978) and was associated with the optical source LS I+61◦303, a young,
rapidly rotating, 10–15 M B0 Ve star (Gregory et al. 1979). A
young pulsar was at first suggested to be responsible for the observed radio emission (Maraschi & Treves 1981), but no pul-sations have ever been detected, despite intensive searches (e.g. Sidoli et al. 2006).
Massi et al.(2017) studied the correlation between the X-ray luminosity and the X-ray spectral slope in LS I+61◦ 303 and
found a good agreement with that of moderate-luminosity BHs. Along with the quasi-periodic oscillations observed from this system both in radio and X-rays (e.g. Nösel et al. 2018), this supports a microquasar scenario for LS I+61◦303. However, in
this case, it is the only known microquasar that exhibits a regu-lar behaviour, does not demonstrate transitions between various spectral states, and lacks a spectral break up to hard γ-rays. A magnetar-like short burst caught from the source supports the identification of the compact object in LS I+61◦303 with a NS
(Barthelmy et al. 2008;Torres et al. 2012).
Radial velocity measurements of the absorption lines of the primary (Casares et al. 2005;Aragona et al. 2009) showed that LS I +61◦ 303 is on an elliptical (e = 0.537 ± 0.34) orbit.
The orbital period of LS I +61◦ 303 was found to be P ≈
26.5 d from radio observations (Gregory 2002). A strong orbital modulation in LS I +61◦ 303 is also observed in the
opti-cal to infrared (Mendelson & Mazeh 1989;Paredes et al. 1994), X-ray (Paredes et al. 1997), hard X-ray (Zhang et al. 2010), and HE/VHE γ-ray (Abdo et al. 2009;Albert et al. 2009) domains. In the optical band, the orbital period signature is evident not only in the broad-band photometry, but also in the spectral prop-erties of the Hα emission line (Zamanov et al. 1999). Because of the uncertainty in the inclination of the system, the nature of the compact object remains unclear, and it can be either a NS or a stellar-mass BH (Casares et al. 2005).
In radio, LS I+61◦ 303 was intensively monitored at GHz
frequencies for many years (e.g. Ray et al. 1997; Massi et al. 2015). The radio light curve displays periodic outbursts whose position and amplitude changed from one orbit to the next. A Bayesian analysis of radio data allowed Gregory (2002) to establish a super-orbital periodic modulation of the phase and amplitude of these outbursts with a period of Pso =
1667 ± 8 days. This modulation has also been observed in X-rays (Chernyakova et al. 2012; Li et al. 2014) and γ-rays (Ackermann et al. 2013;Ahnen et al. 2016;Xing et al. 2017). It has been suggested that the super-orbital periodicity can depend on the Be star disc, either due to a non-axisymmetric struc-ture rotating with a period of 1667 days (Xing et al. 2017), or because of a quasi-cyclic build-up and decay of the Be decretion disc (Negueruela et al. 2001;Ackermann et al. 2013; Chernyakova et al. 2017). Another possible scenario for the super-orbital modulation is related to the precession of the Be star disc (Saha et al. 2016) or periodic Doppler-boosting effects of a precessing jet (Massi & Torricelli-Ciamponi 2016).
The precessing jet model is based on high-resolution radio observations suggesting a double-sided jet (Massi et al. 1993, 2004;Paredes et al. 1998). The precession period in this model is about 26.9 days, which is very close to the orbital period. In this case the observed super-orbital variability is explained as a beat period of the orbital and precession periods (Massi & Jaron 2013).
At GeV energies, LS I +61◦ 303 was unambiguously
detected by Fermi-LAT (Abdo et al. 2009) through its flux mod-ulation at the orbital period. The Fermi-LAT light curve shows a broader peak after periastron and a smaller peak just before apastron (Jaron & Massi 2014). The peak at apastron is affected by the same orbital shift as the radio outbursts and varies on the super-orbital timescale, leading to a decline in the orbital flux modulation as the two peaks merge.
A long-term investigation of Fermi-LAT data bySaha et al. (2016) showed the orbital spectral variability of the source. The observed spectrum is consistent with an exponential cut-off
power law with a cut-off at 6–30 GeV for different orbital states of the system. The excess above the spectral cut-off is part of a second emission component that is dominant at the TeV domain (Hadasch et al. 2012;Saha et al. 2016).
Detected at TeV energies by MAGIC (Albert et al. 2006) and by VERITAS (Acciari et al. 2008), the VHE emission from LS I+61◦ 303 shows a modulation consistent with the orbital
period (Albert et al. 2009) with the flux peaking at apastron. A decade-long VERITAS observation of LS I +61◦ 303 allowed
TeV emission to be detect from the system throughout the entire orbit, with the integral flux above 300 GeV varying in the range (3−7) × 10−12cm−2s−1. The VHE emission is well described
by a simple power-law spectrum, with a photon index of Γ = 2.63 ± 0.06 near apastron andΓ = 2.81 ± 0.16 near periastron (Kar & VERITAS Collaboration 2017).
Similar to other wavelengths, the TeV curve varies from orbit to orbit. MAGIC observations during 2009–2010 caught LS I+61◦303 in a low state, with the TeV flux about an order
of magnitude lower than was previously detected at the same orbital phase (Aleksi´c et al. 2012b).
Long-term multi-wavelength monitoring of LS I+61◦ 303
indicates a correlation between the X-ray (XMM-Newton and Swift/XRT) and TeV (MAGIC and VERITAS) data sets. At the same time, GeV emission shows no correlation with the TeV emission. Along with the spectral cut-off at GeV energies, this implies that the GeV and TeV γ-rays originate from differ-ent particle populations (Anderhub et al. 2009;Aliu et al. 2013; Kar & VERITAS Collaboration 2015).
3.2.2. Prospects for CTA observations
A correlation of VHE and X-ray emission might indicate that in this source the synchrotron emission that is visible at X-rays is due to the same electrons that produce the TeV emission by inverse Compton scattering of stellar photons. However, while X-ray vari-ability on timescales of thousands of seconds is known from the source (Sidoli et al. 2006), MAGIC and VERITAS observations require much longer exposure times, making it difficult to clearly compare the spectral behaviour at different energies. The sensi-tivity of the CTA is crucial for detecting spectral variability on comparable timescales at X-rays and VHE energies.
We studied the capabilities of the CTA to unambigu-ously detect spectral variability of LS I +61◦ 303 on
differ-ent timescales by performing a series of simulations that we based on existing observations. It has been observed that the spectrum of the TeV emission varies between the low and the high state. For our simulations we chose F(E > 1 TeV)= 2.6 × 10−12cm−2s−1 for the high state and F(E > 1 TeV) = 1.2 ×
10−12cm−2s−1for the low state. We assumed a power-law model
to describe the source and used different spectral slopes to deter-mine whether CTA would be able to distinguish between them. The spectral slopes that were chosen wereΓ = 2.4, 2.7, and 3.0, in agreement with MAGIC and VERITAS observations. Simu-lations of both 30 min and 5 h exposure times where performed for each set of parameters. Each combination was simulated 500 times to ensure enough statistics. In the analysis of each real-isation, the normalisation and spectral index were kept as free parameters.
Similarly to the simulations presented in Fig. 2, we found that the uncertainty on the spectral slope has a weak dependence on the slope value. A 5 h observation is enough to determine the slope with an accuracy better than 0.1 (see Figs.1and4).
To determine the capability of the CTA to detect a VHE cut-off in LS I +61◦ 303, we simulated a 5 h observation with a
10−1 100 101 102 E, TeV 10−14 10−13 10−12 10−11 10−10 Flux [erg cm − 2 s − 1]
North 5h high flux
hard mean soft
Fig. 4.Simulated spectra of LS I+61◦303 for 5 h observation. The three
different spectral slopes are shown. Solid lines represent the power-law spectral fit. The butterfly represents the 1σ uncertainty in the fit. power-law spectral model and fit the data with an exponential cut-off power law. The resulting values are shown in Fig.1.
Finally, we studied the orbital variability of the source. We took the light curve above 400 GeV obtained by Albert et al. (2009) and modelled it with 27 bins (see Fig. 5) to study the inter-night variability of the source. We simulated the source with a power-law spectrum with a photon indexΓ = 2.7. In the reconstruction the photon index and the normalisation were left free to vary. We performed 100 realisations for each orbital bin for 30 min and 5 h exposures. In this analysis we assumed a much lower value for the flux in the low state (orbital phases between 0.2 and 0.4) of about 10−14cm−2s−1. The
result-ing uncertainties for the relative flux and slope are summarised in Fig.5. All uncertainties are statistical only at a 1σ confidence level and are below 10% for the integral flux in the high state, with the photon index uncertainty below 0.1 even for a 30 min exposure. In the low state the source is barely detected even with a 5 h exposure. The upper limits we show correspond to the 2σ confidence level. This simulation shows that if the flux of the source is above ≈10−13cm−2s−1, the CTA will be able to detect
inter-night variability of the source at a 10% level. This precision will allow studying the superorbital variability of the orbital pro-file and comparing it to other energy bands. In the high state it will be possible to study the variability of the source at a 30 min timescale. This is comparable to what is observed in X-ray data (see e.g.Chernyakova et al. 2017).
3.3. LS 5039
3.3.1. Source properties
LS 5039 has the shortest orbital period thus far of all known γ-ray binaries (3.9 d, see Table1). It is also known as V497 Sct, based on ROSAT X-ray data.Motch et al.(1997) first reported it as a high-mass X-ray binary. Its peculiar nature as a per-sistent non-thermal radio emitter was soon revealed after the detection of a bright radio counterpart with the Very Large Array (VLA) by Martí et al. (1998). This has anticipated the capability of the system to accelerate electrons to relativis-tic speeds. Follow-up images obtained with very long baseline interferometry (VLBI) resolved the radio emission into elon-gated features, and as a result, LS 5039 was interpreted as a new microquasar system (Paredes et al. 2000). Moreover, at the same time, it was also tentatively associated with the EGRET γ-ray source 3EG J1824−1514. The confirmation of LS 5039 as
0 5 F(E > 400 GeV) [10 − 12 ph cm − 2s − 1] Albert et al. (2009) 10−15 10−14 10−13 10−12 F(E > 1 T eV) [ph cm − 2s − 1] 30m 5h 0.0 0.5 1.0 δ Flux/Flux 30m 5h −0.2 −0.1 0.0 0.1 0.2 0.3 0.4 0.5 0.6 0.7 0.8 0.9 1.0 orbital phase 10−2 10−1 100 δ Γ 30m 5h
Fig. 5.Upper panel: original orbital light curve of LS I+61◦ 303 as
observed by MAGIC (Albert et al. 2009). Next panel: simulated orbital light curve at E > 1 TeV, where the green line shows the simulated flux and the arrows represent 2σ upper limits. Third panel: relative uncer-tainty of the simulated flux. Bottom panel: unceruncer-tainty in the simulated photon index. In all panels the exposure time is shown with colours: blue corresponds to 30 min and red to 5 h.
an unambiguous (>100 GeV) γ-ray source was finally obtained with H.E.S.S. (Aharonian et al. 2005b).
During the 20 yr since its discovery, the physical picture of LS 5039 has generally evolved from the microquasar scenario to a binary system hosting a young non-accreting NS interact-ing with the wind of a massive O-type stellar companion (see e.g.Dubus 2013and references therein). This is strongly sup-ported by VLBI observations of periodic changes in the radio morphology (Moldón et al. 2012), although no radio pulsations have been reported so far.
At different photon energies, the shape of the LS 5039 light curve varies, as confirmed in the most recent multi-wavelength studies using Suzaku, INTEGRAL, COMPTEL, Fermi-LAT, and H.E.S.S. data (Chang et al. 2016, and references therein). The X-ray, soft γ-ray (up to 70 MeV), and TeV emission peak around inferior conjunction after the apastron passage. In contrast, γ-rays in the 0.1–3 GeV energy range anti-correlate and have a peak near the superior conjunction soon after the periastron pas-sage. No clear orbital modulation is apparent in the 3–20 GeV band. This dichotomy suggests a highly relativistic particle pop-ulation that accounts for both X-ray/soft γ-ray and TeV emission mainly by synchrotron and anisotropic inverse Compton (IC) scattering of stellar photons, respectively. The GeV γ-ray peak would arise when TeV photons (of an IC origin) are absorbed through pair production as the NS approaches its O-type com-panion, and further enhances the GeV emission through cas-cading effects. Variable adiabatic cooling and Doppler boosting are other effects proposed to play an important role when trying to understand the multi-wavelength modulation of systems such as LS 5039 (see e.g.Khangulyan et al. 2008a; Takahashi et al. 2009;Dubus 2013).
3.3.2. Prospects for CTA observations
In order to estimate the capabilities of the CTA of detecting the temporal and spectral variations of emission from LS 5039, we have simulated CTA observations of this source at different spec-tral states. Because the emission spectrum of LS 5039 varies
0.0 0.5 1.0 1.5 2.0 0.5 1.0 1.5 Flux [10 − 12 ph /(cm 2 sec)] Tobs= 0.5 hr Tobs= 5 hr 0.0 0.5 1.0 1.5 2.0 0.1 0.2 ∆ Γ Tobs= 0.5 hr Tobs= 5 hr 0.0 0.5 1.0 1.5 2.0 Orbital phase 10−1 100 101 102 ∆ Ecut [T eV] Tobs= 0.5 hr Tobs= 5 hr
Fig. 6.Simulated CTA view of LS 5039 orbital variations above 1 TeV
energy for 0.5 and 5 h long snapshots as observed from the southern CTA site. Upper panel: integral flux above 1 TeV, and middle and lower panels: estimated uncertainty of the measured spectral index and cut-off energy, correspondingly. The data points corresponding to Tobs= 5 h are
shifted to the right by 0.01 phase for clarity.
with orbital phase, we assumed the flux and spectral shape mod-ulations found in the recent H.E.S.S. data (Mariaud et al. 2015). These observations suggest that the source spectrum varies from dN/dE ∝ E−1.9exp (−E/6.6 TeV) at inferior conjunction (phase
≈0.7) to dN/dE ∝ E−2.4at superior conjunction (phase ≈0.05). We furthermore assumed that the source spectrum always fol-lows a power law with an exponential cut-off shape and took the flux and spectral index evolution from Figs. 3 and 4 of Mariaud et al.(2015). Because no spectral cut-off was observed at superior conjunction, we assumed that the cut-off energy is modulated between Emin
cut = 6.6 TeV at phase φ ≈ 0.7 and
Emaxcut = 40 TeV at phase φ ≈ 0.3:
log10(Ecut)= log10(Emean) −∆ log10E ×cos(φ − 0.71), (1)
where log10(Emean) = 0.5 × [log10(Emaxcut )+ log10(Ecutmin)] and
∆ log10E= 0.5 × [log10(Ecutmax) − log10(Ecutmin)].
We simulated ten snapshot observations from orbital phases 0.0–1.0, lasting 0.5 h and 5 h each. This gives a total exposure of 5 and 50 h on the source. To reconstruct the simulated flux, we assumed the same power law with exponential cut-off model, but with the spectrum normalisation, index, and cutoff energy as free parameters. For each phase bin and observation duration, the simulation was repeated 100 times to estimate the mean values of the flux (in the 1–100 TeV range), spectral index, and cut-off energy, as well as their standard deviations. The results of these simulations are shown in Fig. 6; the estimated uncertainties of the reconstructed source spectral parameters are also given there.
Figure6shows that the CTA can follow the orbital flux evo-lution of LS 5039 even with 30 min observational snapshots. However, in the orbital phase range 0.1−0.3, the uncertainties on the flux become &20% and at least 5 h long exposures would be required to determine the flux accurately. Such exposure times yield .10% accuracy of the flux and spectral index determina-tion in this phase range, whereas the cut-off energy cannot be measured accurately. Only during the bright flux period, corre-sponding to the phase range ≈0.4−0.9, is the uncertainty in the cut-off energy better than ≈20%. This implies that detailed spec-tral studies of this particular binary phase will require integration over several orbital periods.
Such CTA observations of LS 5039 during its high-flux peri-ods will enable spectral studies on timescales as short as 0.01 orbital periods (≈1 h). They will strongly constrain the physical processes at work.
Furthermore, it will clarify whether the rotating hollow-cone model (Neronov & Chernyakova 2008) that has previously been proposed to explain the LS 5039 TeV light curve is feasi-ble. In this model, the TeV peak is composed of two narrower peaks, whose appearance depends on the location of the emis-sion region and the system geometry.Neronov & Chernyakova (2008) suggested that the flux difference between the peaks and inferior conjunction at phase ≈0.7 is &10%. Such variations are detectable with CTA in several hours of exposure, as the top panel of Fig.6shows. In this way, CTA observations may allow constraining the geometry of the system, including the otherwise elusive orbital inclination angle.
3.4. 1FGL J1018.6−5856 3.4.1. Source properties
1FGL J1018.6−5856 (3FGL J1018.9−5856, HESS J1018−589A) is a point like γ-ray source. It is positionally coincident with the supernova remnant SNR G284.3–1.8. Using the Gaia DR2 source parallax and assuming a Gaussian probability distribu-tion for the parallax measurement,Marcote et al.(2018) derived a source distance of d= 6.4+1.7
−0.7kpc. They also calculated the
Galactic proper motion of the source and found that it is mov-ing away from the Galactic plane. Both the source distance and proper motion are not compatible with the position of the SNR G284.3−1.8 (which is located at an estimated distance of '2.9 kpc). Therefore, it is possible to exclude any physical rela-tion between the binary source and the SNR.
Spectroscopic observations of the optical counterpart allowedStrader et al. (2015) to find that a companion star has a low radial velocity semi-amplitude of 11–12 km s−1, which
favours a NS as a compact object. This conclusion is in agree-ment with the results of Monageng et al. (2017), who con-strained the eccentricity of the orbit e= 0.31 ± 0.16 and showed that the compact object is a NS, unless the system has a low inclination i . 26◦.
The 1FGL J1018.6–5856 was detected in a blind search for periodic sources in the Fermi-LAT survey of the Galactic Plane. Optical observations show that the non-thermal source is posi-tionally coincident with a massive star of spectral type O6V(f). The radio and X-ray fluxes from the source are modulated with the same period of 16.544 days, interpreted as the binary orbital period (Fermi LAT Collaboration 2012).
The very high-energy counterpart of this source is the point-like source HESS J1018−589A, (H.E.S.S. Collaboration 2012a). In a dedicated observation campaign at VHE, HESS J1018−589A was detected up to 20 TeV. Its energy spectrum is well described
with a power-law model, with a photon indexΓ = 2.2 and a mean differential flux N0 = (2.9 ± 0.4) × 10−13ph cm−2s−1TeV−1at
1 TeV. As in the case of other γ-ray binaries, the VHE spectrum cannot be extrapolated from the HE spectrum, which has a break at around 1 GeV. The orbital light curve at VHE peaks in phase with the X-ray and HE (1–10 GeV) light curves.
Based on optical spectroscopic observations, Strader et al. (2015) found that the maxima of the X-ray, HE, and VHE flux correspond to the inferior conjunction. This finding was unexpected because γ-rays are believed to be produced through anisotropic inverse Compton up-scattering of the stellar UV pho-tons. Therefore, the peak of the γ-ray flux should occur at the superior conjunction, especially if the system is edge-on. This discrepancy could only be explained if the binary orbit is eccen-tric and the flux maximum occurs at periastron.
NuSTAR observations (An et al. 2015) demonstrated that similar to other γ-ray binaries, the broad-band X-ray spectrum is well fitted with an unbroken power-law model. The source flux shows a correlation with the spectral hardness throughout all orbital phases.
A comparison of the light curves of 1FGL J1018.6−5856 at different energy ranges shows that both the X-ray and the low-energy (E < 0.4 GeV) γ-ray bands are characterised by a similar modulation (a broad maximum at φ= 0.2–0.7 and a sharp spike at φ= 0), thus suggesting that they are due to a common spectral component. On the other hand, above ≈1 GeV, the orbital light curve changes significantly because the broad hump disappears and the remaining structure is similar to the light curve observed at VHE. Based on these results,An & Romani(2017) suggested that the flux in the GeV band is due mainly to the pulsar mag-netosphere, while the X-ray flux is due to synchrotron emis-sion from shock-accelerated electrons and the TeV light curve is dominated by the up-scattering of the stellar and synchrotron photons through external Compton (EC) and synchrotron self-Compton (SSC) mechanisms, in an intrabinary shock. The light curves at different energy ranges can be reproduced with the beamed SSC radiation from adiabatically accelerated plasma in the shocked pulsar wind. This is composed of a slow and a fast outflow. Both components contribute to the synchrotron emis-sion observed from the X-ray to the low-energy γ-ray band, which has a sinusoidal modulation with a broad peak around the orbit periastron at φ= 0.4. On the other hand, only the Doppler-boosted component reaches energies above 1 GeV, which are characterised by the sharp maximum that occurs at the inferior conjunction at φ= 0. This result can be obtained with an orbital inclination of ≈50◦and an orbital eccentricity of ≈0.35,
consis-tent with the constraints obtained from optical observations. In this way, the model could also explain the variable X-ray spike coincident with the γ-ray maximum at φ= 0.
3.4.2. Prospects for CTA observations
Although 1FGL J1018.6−5856 was investigated in depth over the past few years, several question about its properties are still open, such as the physical processes that produce the HE/VHE emission. Moreover, it is still not clear whether the X-ray and γ-ray peaks are physically related to the conjunctions or the apastron and periastron passages. Therefore, the observation of this source with CTA will allow us to address a few topics. The high sensitivity of CTA will enable us to investigate the orbital modulation of the source spectrum and to study the correla-tion of the VHE emission with the system geometry. From the spectral point of view, the spectral shape will be further con-strained at both the low (E < 0.1 TeV) and the high (E > 20 TeV)
0.110 1 10 100 −13 10 −12 2×10 −13 5×10 −13 E 2 x dN/dE (erg cm −2 s −1) E, TeV 0.110 1 10 100 −13 10 −12 2×10 −13 5×10 −13 E 2 x dN/dE (erg cm −2 s −1) E, TeV
Fig. 7. Real and simulated spectra of 1FGL J1018.6−5856. Black:
spectrum of 1FGL J1018.6−5856 obtained with 63 h of observation with H.E.S.S. (H.E.S.S. Collaboration 2015a). Green: simulation of the source spectrum obtained with 50 h of observation with CTA.
energy end. This will provide further constraints on the location, magnetic field, and acceleration efficiency of the VHE emitter (Khangulyan et al. 2008b) and on the opacity due to pair pro-duction (Böttcher & Dermer 2005;Dubus 2006).
To study the CTA capabilities, we first simulated the phase-averaged spectrum of the source based on the H.E.S.S. obser-vations (H.E.S.S. Collaboration 2015a). As input we assumed a simple power-law emission with a photon indexΓ = 2.2 and a flux normalisation N0 = 2.9 × 10−19MeV−1cm−2s−1 at 1 TeV.
We performed three sets of simulations for 30 min, 5 h, and 50 h observations. For each set we performed 100 simulations. Figure1shows that with a 5 h observation, it will be possible to measure the source flux and spectral slope with an uncertainty of '10% and 0.05, respectively. With only 30 min of observation, the corresponding errors would be '35% and 0.3.
In Fig. 7 we report the simulated spectrum of 1FGL J1018.6−5856 obtained with 50 h of observation, together with that obtained with 63 h of H.E.S.S. observations. It shows that the CTA spectrum is well determined both at low (down to E ' 0.1 TeV) and high energies (up to E ' 100 TeV), thus provid-ing a significant enlargement of the spectral coverage compared to H.E.S.S.. The extension of the spectral range towards low energies will enable investigating the connection to the MeV– GeV emission, while the increase of the high-energy end will be important to constrain the cut-off linked to particle acceleration. We estimate that with a 5 h observation, the CTA will allow us to detect a high-energy cut-off if it is located below 18 TeV (Fig.1). We also studied the flux modulation of the source throughout the orbit. FollowingH.E.S.S. Collaboration(2015a), we divided the whole orbit into ten phase bins (1 bin ' 39.67 h) and assumed a simple photon spectrum with a power-law model with a slope Γ = 2.2. For each phase bin we performed 100 simulations for both 30min and 5h observations. We fitted the spectrum with a simple power-law model, keeping both the normalisation and the photon indexΓ free to vary. The results of this set of simulations are reported in Fig.8, where we show the variability (throughout the orbital phase) of the flux, its relative error, and the uncer-tainty of the index.
In the upper panel of Fig.8, we report with red symbols the simulated light curve obtained with a campaign of 5 h observa-tions for each phase bin with the CTA. It proves that in this case, the CTA can clearly resolve the source flux variability through-out the orbit. It will be possible to point through-out flux variations of
0 1 2 3 Flux HESS 0 2 4 6 Flux CTA 0 0.2 0.4 0.6 dF/F 0 0.2 0.4 0.6 0.8 1 0 0.2 0.4 0.6 d Γ Phase
Fig. 8.Simulated CTA view of 1FGL J1018.6−5856 orbital
variabil-ity. Upper panel: flux modulation of 1FGL J1018.6−5856 throughout the orbital phase as observed with H.E.S.S. (black symbols, in units of 10−12ph cm−2s−1for E > 0.35 TeV) and simulated for 5 h of observation
with the CTA (red symbols, in units of 10−13ph cm−2s−1for E > 1 TeV).
Middle panel: flux relative error in the case of simulated CTA observa-tions of 30 min (blue filled squares) and 5 h (red open squares). Lower panel: uncertainty of the photon indexΓ in the same cases.
'25% over timescales of ≈0.1 orbital periods (middle panel), and even at its flux minimum, the source will be detected with a significance &7σ. For comparison, in the upper panel of Fig.8 we report as black symbols the folded light curve obtained with H.E.S.S. (with '8 h of observation for each phase bin). We note that in this case, it is possible to claim a clear flux variability only in the three phase bins in the phase range φ= 0.8–1.1, while the flux values measured in the remaining phase range are consis-tent with each other. The better characterisation of the source variability provided by the CTA will enable an improved corre-lation with the X-ray and HE variability and will place tighter constraints on the position and size of the VHE emitter.
The lower panel of Fig.8shows that even at the flux mini-mum, 5 h of observation with the CTA will provide a measure-ment of the photon index with an uncertainty of 0.2 ('9%). This accuracy is comparable to that obtained with more than 60 h of observation with H.E.S.S. Therefore, it will be feasible to deter-mine possible spectral variations >10% in the orbital phase. In this way, it will be possible to single out the VHE emission and absorption processes and to obtain useful information on both the source magnetic field and the efficiency of the particle accel-eration and pair production.
3.5. HESS J0632+057 3.5.1. Source properties
In contrast to other γ-ray loud binaries, HESS J0632+057 remained the only system that for a time was lacking detec-tion in the GeV energy band. Only recently have indica-tions of a GeV detection with Fermi-LAT been reported by Malyshev & Chernyakova(2016) andLi et al.(2017). The sys-tem was initially discovered during H.E.S.S. observations of the Monoceros region (Aharonian et al. 2007) as an uniden-tified point-like source. Its spatial coincidence with the Be star MWC 148 suggested its binary nature (Aharonian et al. 2007; Hinton et al. 2009). With dedicated observational cam-paigns, the binary nature of the system was confirmed by radio (Skilton et al. 2009) and soft X-ray (Falcone et al. 2010)
observations. In the TeV band, the system was also detected by VERITAS and MAGIC (Aleksi´c et al. 2012c;Aliu et al. 2014).
The orbital period of HESS J0632+057 of ≈316 ± 2 d (Malyshev et al. 2017), with a zero-phase time T0= 54857 MJD
(Bongiorno et al. 2011), was derived from Swift/XRT observa-tions. The exact orbital solution and even the orbital phase of periastron is not firmly established and is placed at orbital phases φ ≈ 0.97 (Casares et al. 2012) or φ ≈ 0.4−0.5 (Moritani et al. 2018;Malyshev et al. 2017).
The orbital folded X-ray light curve of HESS J0632+057 has two clear emission peaks: first at phase φ ≈ 0.2−0.4, and second at φ ≈ 0.6−0.8 separated by a deep minimum at φ ≈ 0.4−0.5 (Bongiorno et al. 2011; Aliu et al. 2014). A low-intermediate state is present at φ ≈ 0.8−0.2. The orbital light curve in the TeV energy range shows a similar structure, as was reported byMaier & VERITAS Collaboration(2015). Indi-cations of orbital variability in the GeV range were reported byLi et al.(2017).
The X-ray-to-TeV spectrum of HESS J0632+057 is shown in Fig.9. Several models have been proposed so far to explain the observed variations of the flux and spectrum throughout the orbit. In the flip-flop scenario (see e.g.Moritani et al. 2015, and references therein) the compact object is assumed to be a pulsar that passes periastron at φ = 0.97. Close to apastron (orbital phases ≈0.4P−0.6), the pulsar is in a rotationally pow-ered regime, while it switches into a propeller regime when peri-astron is approached (phases 0.1−0.4 and 0.6−0.85). When the gas pressure of the Be disc overcomes the pulsar-wind ram pres-sure, the pulsar wind in a flip-flop scenario is quenched (phases 0−0.1 and 0.85−1). Because the Be disc of the system is esti-mated to be about three times larger than the binary separation at periastron, the compact object enters a dense region of the disc near periastron. In this situation, the strong gas pressure is likely to quench the pulsar wind and suppress high-energy emissions. Alternatively, the observed orbital variations can be explained within the “similar to PSR B1259−63” model (Malyshev et al. 2017). The similar two-peak behaviour of the HESS J0632+057 and PSR B1259−63 orbital light curves allows us to assume that the orbital plane of HESS J0632+057 is inclined with respect to the disc plane, similarly to PSR B1259−63. Orbital X-ray and TeV peaks within this model correspond to the first and second crossing of the disc by a compact object. Higher ambient den-sity during these episodes leads to more effective cooling of the relativistic electrons by synchrotron and inverse Compton mech-anisms, resulting in an increased level of X-ray and TeV emis-sion. The orbital phase of periastron in this model is located at phase φ ≈ 0.4−0.5 (Malyshev et al. 2017).
The break in the GeV–TeV spectrum at ≈200 GeV can be interpreted as a corresponding break in the spectrum of emit-ting relativistic electrons. The X-ray-to-GeV and TeV parts of the spectrum are explained as synchrotron and IC components. An initial power-law (Γ1,e ≈ 1.3) spectrum of electrons can be
modified by synchrotron energy losses at above Ebr ≈ 1 TeV,
resulting in aΓ2,e ≈ 2.3 higher energy slope. The absence of
cooling in the energy band below 1 TeV could be attributed to the escape of the sub-TeV electrons from the system. A similar interpretation of the spectral energy distribution was proposed byChernyakova et al.(2015) for PSR B1259−63.
Alternatively, the spectral break in the electron spectrum can occur at the transition between the domination of adiabatic and IC or synchrotron losses (see e.g. Khangulyan et al. 2007 and Takahashi et al. 2009for PSR B1259−63 and LS 5039). The adi-abatic loss time is naturally shortest in sparse regions outside of the Be star disc and longest in dense regions inside it.
10
410
610
810
1010
12E, eV
10
1410
1310
1210
11EF
E,
er
g/
cm
2/s
Fig. 9. X-ray-to-TeV spectrum of HESS J0632+057 during its high
state (green points; orbital phases φ ≈ 0.3−0.4) and low state (blue points; φ ≈ 0.4−0.5). The data are adopted from Malyshev et al.
(2017) (X-rays),Li et al.(2017) (mean GeV spectrum, black points), and Malyshev & Chernyakova(2016) (green upper limits). TeV data are adopted fromMaier & VERITAS Collaboration(2015). The solid lines show the “similar to PSR B1259−63” model flux, while dashed and dot-dashed lines illustrate contributions from synchrotron and IC model components correspondingly. See text for more details.
A broken power-law shape of the spectrum is not unique for the “similar to PSR B1259−63” model. A similar shape of the spectrum can also be expected within the flip-flop model because both interpretations of the break origin can be valid for this model. The two models can be distinguished by CTA obser-vations of the variation in slope and low-energy break position throughout the orbit.
Within the flip-flop model at orbital phases φ = 0−0.4, the compact object moves from a denser to increasingly sparser regions of the Be star disc. The spectrum of relativistic electrons becomes increasingly less dominated by the losses. This results in a gradual hardening of the TeV slope and in a shift of the break energy to higher values. At phases φ ≈ 0.6−1, the com-pact object enters denser regions of the disc, which should lead to a gradual softening of the slope and shift of the energy break to lower energies. The spectrum is expected to be hardest when the object is beyond the Be star disc (orbital phase ≈0.4). This phase corresponds to the minima of observed emission. The soft-est spectrum is expected when the compact object approaches periastron, that is, at phase φ ≈ 0.97.
In the “similar to PSR B1259−63” model the compact object intersects the disc of the Be star twice per orbit (at orbital phases 0.2−0.4 and 0.6−0.8) where the soft spectrum with the low posi-tion of energy break is expected. At phases 0−0.2, 0.4−0.6 and 0.8−1 in the “similar to PSR B1259−63” model, the compact object is beyond the dense regions of the disc. At these orbital phases a hard slope with energy break shifted to higher energies can be expected.
3.5.2. Prospects for CTA observations
Because of its location, HESS J0632+057 is visible from both the north and south CTA sites (see Table 1). For our simula-tions we considered two orbital phases: the brightest phase (φ= 0.2–0.4, hereafter the “high state”) based onAliu et al.(2014), and the low-intermediate phase (φ = 0.8–0.2, hereafter the
“low state”) based on Schlenstedt (2017). No spectra are reported in the literature for the deep minimum state at φ= 0.4– 0.5 and the two maxima have similar spectra, therefore we chose the brightest as representative of the active state.
In the first group of simulations, we considered the 0.1– 100 TeV energy range. The spectral model component of the source was defined as a power-law model with either a photon index 2.3 and normalisation at 1 TeV of 5.7 × 10−13ph cm−2s−1
TeV−1 (φ = 0.2–0.4, high state), or a photon index 2.72 and
normalisation 2.3 × 10−13ph cm−2s−1TeV−1 (φ = 0.8–0.2, low
state).
The dependence of the source flux and index uncertain-ties on different configurations is shown in Fig. 1 for 30 min and 5 h observations. Simulations also show that a longer 50 h observation can reconstruct the flux and slope of the source to an accuracy of better than 3% in the Northern and Southern Hemispheres.
In the second group of simulations, we simulated the source spectrum with a broken power-law model to study the possi-bility of the detection of a low-energy break. The two phys-ical scenarios discussed above, flip-flop versus “similar to PSR B1259−63”, can be distinguished by CTA observations of the variation in the position of the slope and low-energy break throughout the orbit. In both high and low states, we used a low-energy slope fixed at 1.6 and an energy break at 0.4 TeV (the value of the break energy was not fixed in the simulations, therefore it can vary between different realisa-tions,Malyshev & Chernyakova 2016). The high-energy slopes are given by the spectral indices already discussed: Γ = 2.3 for the φ = 0.2–0.4 high state, and Γ = 2.72 for the φ = 0.8–0.2 low state. The normalisations for the two states at 0.4 TeV are 4.7 × 10−12ph cm−2s−1TeV−1(high state), and 2.8 ×
10−12ph cm−2s−1TeV−1 (low state). For these simulations, we
focused on the southern site and the 0.04–100 TeV energy range. Simulations of 5 h and 50 h observations were performed, and the results are shown in Table2.
Figure 10 shows two single realizations of the spectra of HESS J0632+057, 5 h (left) versus 50 h (right). Upper limits are shown when the detection significance is lower than 3σ. A 50 h observation will excellently reconstruct the energy break and slopes, whereas a 5 h observation will suffer from higher uncertainties.
A direct comparison of the same power-law spectra simu-lated for the CTA with respect to H.E.S.S. (Aliu et al. 2014) and VERITAS (Schlenstedt 2017) observations shows that a CTA snapshot of 5 h will result in more accurate results than what was previously obtained. A 55 h observation of the low state (φ = 0.8–0.2) of the source with VERITAS resulted in a ≈7% uncertainty on the detected slope (2.72 ± 0.2), to be compared to the ≈5% with a 5 h CTA south observation. Similarly, a 15 h observation of the high state with H.E.S.S. resulted in a ≈9% uncertainty on the detected slope (2.3 ± 0.2), to be compared to the ≈3% with a 5 h CTA south observation. The CTA error esti-mates given here are purely statistic, whereas the VERITAS and H.E.S.S. results include systematic errors as well.
Figure 1 and Table 2 show that a 5 h observation will be enough to distinguish the low state from the high state, but it may not be enough to unambiguously distinguish the energy break (≈30% uncertainty). A 50 h observation would result in a 10% uncertainty of the energy break (15% for a 20 h obser-vation), allowing the high-energy slope and energy break to be accurately monitored throughout the orbit. This would enable the CTA to distinguish the two currently available scenarios, that is, flip-flop versus “similar to PSR B1259−63”, which expect an
Table 2.Best fit of the simulated spectra shown in Fig.10(0.04–100 TeV) with a broken power-law model for HESS J0632+057.
5 h 50 h
Phase Eb Γ Flux Eb Γ Flux
0.2–0.4 0.51 ± 0.10 2.30 ± 0.07 1.09 ± 0.16 0.40 ± 0.04 2.30 ± 0.02 1.07 ± 0.05 0.8–0.2 0.46 ± 0.14 2.73 ± 0.14 0.65 ± 0.15 0.40 ± 0.04 2.71 ± 0.05 0.62 ± 0.04
Notes. Eb(TeV) is the position of energy break,Γ is the photon index above the break (the low-energy photon index was frozen to 1.6), and F is
the 0.04–100 TeV flux in 10−11ph cm−2s−1units. See text for more details.
101 100 101 102 E, TeV 1013 1012 EFE ,e rg /cm 2/s 101 100 101 102 E, TeV 1013 1012 EFE ,e rg /cm 2/s 101 100 101 102 E, TeV 1014 1013 1012 EFE ,e rg /cm 2/s 101 100 101 102 E, TeV 1014 1013 EFE ,e rg /cm 2/s
Fig. 10.Simulated spectra (red and green points) of HESS J0632+057, as observed from the southern site. In blue we show the input models.
Upper left: high state, 5 h. Upper right: high state, 50 h. Lower left: low state, 5 h. Lower right: low state, 50 h. opposite trend of the spectral slope and energy break from the
high state to the low state: a hardening of the spectrum and Eb
moving to higher energies for the “similar to PSR B1259−63” scenario versus a softening of the spectrum and Eb moving to
lower energies in the flip-flop model.
HESS J0632+057 has a long orbital period (≈316 d), and each phase will occur only once in one year of observations. Nevertheless, each state is observable for a long period, there-fore a 50 h observation in the same state (10 nights with ≈5 h each) is possible.
3.6. HESS J1832−093 3.6.1. Source properties
HESS J1832−093 is a new γ-ray binary candidate discov-ered as a TeV point source by H.E.S.S. This source lies in
the vicinity of SNR G22.7−0.2, which might suggest an asso-ciation with this SNR (H.E.S.S. Collaboration 2015b). How-ever, several follow-up observations in X-rays instead support the binary nature of this source (Eger et al. 2016; Mori et al. 2017). A simple power-law model describes the TeV spec-trum well, with a photon index ofΓ = 2.6 ± 0.3stat ± 0.1sys and an integrated photon flux above 1 TeV of F= (3.0 ± 0.8stat ± 0.6syst) × 10−13cm−2s−1 (H.E.S.S. Collaboration 2015b). An
XMM-Newtonobservation of the source field discovered a bright X-ray source, XMMU J183245−0921539, within the γ-ray error circle (H.E.S.S. Collaboration 2015b). This source is also associated with a point source detected in a subsequent Chan-draobservation campaign (Eger et al. 2016). During the Chan-draobservations, an increase of the 2–10 keV flux of the order of 4 with respect to the earlier XMM-Newton measurement and the coincidence of a bright IR source at the Chandra error box sug-gest a binary scenario for the γ-ray emission (Eger et al. 2016).
